The Low-Mass Populations in OB Associations

Home , IC 348, Melnick 34, NGC 206

Briceño et al.: Low-mass Populations in OB Associations 345

César Briceño Centro de Investigaciones de Astronomía

Thomas Preibisch Max-Planck-Institut für Radioastronomie

William H. Sherry National Optical Astronomy Observatory

Eric E. Mamajek Harvard-Smithsonian Center for Astrophysics

Robert D. Mathieu University of Wisconsin–Madison

Frederick M. Walter Stony Brook University

Hans Zinnecker Astrophysikalisches Institut Potsdam

Low-mass stars (0.1 < M < 1 M ) in OB associations are key to addressing some of the most fundamental problems in star formation. The low-mass stellar populations of OB associations provide a snapshot of the fossil star-formation record of giant molecular cloud complexes. Large-scale surveys have identified hundreds of members of nearby OB associations, and revealed that low-mass stars exist wherever high-mass stars have recently formed. The spatial distribution of low-mass members of OB associations demonstrate the existence of significant substructure (“subgroups”). This “discretized” sequence of stellar groups is consistent with an origin in short-lived parent molecular clouds within a giant molecular cloud complex. The low- mass population in each subgroup within an OB association exhibits little evidence for significant age spreads on timescales of ~10 m.y. or greater, in agreement with a scenario of rapid star formation and cloud dissipation. The initial mass function (IMF) of the stellar populations in OB associations in the mass range 0.1 < M < 1 M is largely consistent with the field IMF, and most low-mass pre-main-sequence stars in the solar vicinity are in OB associations. These findings agree with early suggestions that the majority of stars in the galaxy were born in OB associations. The most recent work further suggests that a significant fraction of the stellar population may have their origin in the more spread out regions of OB associations, instead of all being born in dense clusters. Groundbased and spacebased (Spitzer Space Telescope) infrared studies have provided robust evidence that primordial accretion disks around low-mass stars dissipate on timescales of a few million years. However, on close inspection there appears to be great variance in the disk dissipation timescales for stars of a given mass in OB associations. While some stars appear to lack disks at ~1 m.y., a few appear to retain accretion disks up to ages of ~10–20 m.y.

1. INTRODUCTION expanding stellar systems of blue luminous stars. These generally include groups of T Tauri stars or T associations Most star formation in normal galaxies occurs in the (Kholopov, 1959; Herbig, 1962; Strom et al., 1975) as well cores of the largest dark clouds in spiral arms, known as as clusters, some containing massive (M > 10 M ) stars, but giant molecular clouds (GMCs). A GMC may give rise to all teeming with solar-like and lower-mass stars. one or more star complexes known as OB associations, first Although we now recognize OB associations as the prime defined and recognized by Ambartsumian (1947) as young sites for star formation in our galaxy, much of our knowl-

345 346 Protostars and Planets V edge of star formation is based on studies of low-mass (M < 3. Bound vs. unbound clusters. While many young 1 M ) pre-main-sequence (PMS) stars located in nearby T stars are born in groups and clusters, most disperse rapidly; associations, like the ~1–2 m.y. old Taurus, Lupus, and Cha- few clusters remain bound over timescales >10 m.y. The maeleon star-forming regions. The view of star formation conditions under which bound clusters are produced are not conveyed by these observations is probably biased to the clear. Studies of older, widely spread low-mass stars around particular physical conditions found in these young, quies- young clusters might show a time sequence of cluster for- cent regions. In contrast, the various OB associations in the mation, and observations of older, spreading groups would solar vicinity are in a variety of evolutionary stages and yield insight into how and why clusters disperse. environments, some containing very young objects (ages 4. Slow vs. rapid disk evolution. Early studies of near- <1 m.y.) still embedded in their natal gas (e.g., Orion A and infrared dust emission from low-mass young stars suggested B clouds, Cep OB2), others in the process of dispersing their that most stars lose their optically thick disks over periods parent clouds, like λ Ori and Carina, while others harbor of ~10 m.y. (e.g., Strom et al., 1993), similar to the time- more evolved populations, several million years old, that scale suggested for planet formation (Podosek and Cassen, have long since dissipated their progenitor clouds (like Scor- 1994). However, there is also evidence for faster evolution pius-Centaurus and Orion OB1a). The low-mass popula- in some cases; for example, half of all ~1-m.y.-old stars in tions in these differing regions are key to investigating fun- Taurus have strongly reduced or absent disk emission (Beck- damental issues in the formation and early evolution of stars with et al., 1990). The most recent observations of IR emis- and planetary systems: sion from low-mass PMS stars in nearby OB associations 1. Slow vs. rapid protostellar cloud collapse and mo- like Orion suggest that the timescales for the dissipation of lecular cloud lifetimes. In the old model of star forma- the inner disks can vary even in coeval populations at young tion (see Shu et al., 1987) protostellar clouds contract slowly ages (Muzerolle et al., 2005). until ambipolar diffusion removes enough magnetic flux for 5. Triggered vs. independent star formation. Although dynamical (inside-out) collapse to set in. It was expected it is likely that star formation in one region can “trigger” that the diffusion timescale of ~10 m.y. should produce a more star formation later in neighboring areas, and there is similar age spread in the resulting populations of stars, evidence for this from studies of the massive stars in OB consistent with the <40 m.y. early estimates of molecular populations (e.g., Brown, 1996), proof of causality and cloud lifetimes (see discussion in Elmegreen, 1990). Such precise time sequences are difficult to obtain without study- age spreads should be readily apparent in color-magnitude ing the associated lower-mass populations. In the past, stud- or HR diagrams for masses <1 M . However, the lack of ies of the massive O and B stars have been used to investi- even ~10-m.y.-old, low-mass stars in and near molecular gate sequential star formation and triggering on large scales clouds challenged this paradigm, suggesting that star forma- (e.g., Blaauw, 1964, 1991, and references therein). How- tion proceeds much more rapidly than previously thought, ever, OB stars are formed essentially on the main sequence even over regions as large as 10 pc in size (Ballesteros- (e.g., Palla and Stahler, 1992, 1993) and evolve off the main Paredes et al., 1999), and therefore that cloud lifetimes over sequence on a timescale on the order of 10 m.y. (depend- the same scales could be much shorter than 40 m.y. (Hart- ing upon mass and amount of convective overshoot), thus mann et al., 1991). they are not useful tracers of star-forming histories on time- 2. The shape of the IMF. Whether OB associations scales of several million years, while young low-mass stars have low-mass populations according to the field IMF, or are. Moreover, we cannot investigate cluster structure and if their IMF is truncated, is still a debated issue. There have dispersal or disk evolution without studying low-mass stars. been many claims for IMF cutoffs in high-mass star-form- Many young individual clusters have been studied at both ing regions (see, e.g., Slawson and Landstreet, 1992; optical and infrared wavelengths (cf. Lada and Lada, 2003), Leitherer, 1998; Smith et al., 2001; Stolte et al., 2005). How- but these only represent the highest-density regions, and do ever, several well-investigated massive star-forming regions not address older and/or more widely dispersed popula- show no evidence for an IMF cutoff [see Brandl et al. (1999) tions. In contrast to their high-mass counterparts, low-mass and Brandner et al. (2001) for the cases of NGC 3603 and stars offer distinct advantages for addressing the aforemen- 30 Dor, respectively], and notorious difficulties in IMF de- tioned issues. They are simply vastly more numerous than terminations of distant regions may easily lead to wrong con- O, B, and A stars, allowing statistical studies not possible clusions about IMF variations (e.g., Zinnecker et al., 1993; with the few massive stars in each region. Their spatial dis- Selman and Melnick, 2005). An empirical proof of a field- tribution is a fossil imprint of recently completed star for- like IMF, rather than a truncated IMF, has important con- mation, providing much needed constraints for models of sequences not only for star-formation models but also for molecular cloud and cluster formation and dissipation; with scenarios of distant starburst regions; e.g., since most of the velocity dispersions of ~1 km s–1 (e.g., de Bruijne, 1999) stellar mass is then in low-mass stars, this limits the amount the stars simply have not traveled far from their birth sites of material that is enriched in metals via nucleosynthesis in (~10 pc in 10 m.y.). Low-mass stars also provide better massive stars and that is then injected back into the inter- kinematics, because it is easier to obtain accurate radial stellar medium by the winds and supernovae of the mas- velocities from the many metallic lines in G-, K-, and M- sive stars. type stars than it is from O- and B-type stars. Briceño et al.: Low-mass Populations in OB Associations 347

2. SEARCHES FOR LOW-MASS spectral types and a qualitative assessment of the strength PRE-MAIN-SEQUENCE STARS of prominent emission lines, like the hydrogen Balmer lines IN OB ASSOCIATIONS or the Ca II H and K lines. Liu et al. (1981) explored a 5° × 5° region in Per OB2 and detected 25 candidate TTS. Ogura Except for the youngest, mostly embedded populations in (1984) used the 1-m Kiso Schmidt to find 135 Hα-emitting the molecular clouds, or dense, optically visible clusters like stars in Mon OB1. Wilking et al. (1987) detected 86 emis- the Orion Nebula Cluster (ONC), most of the low-mass stel- sion line objects over 40 deg2 in the ρ Ophiuchi complex. lar population in nearby OB associations is widely spread Mikami and Ogura (2001) searched an area of 36 deg2 in over tens or even hundreds of square degrees on the sky. Cep OB3 and identified 68 new emission line sources. In Moreover, it is likely that after ~4 m.y. the stars are no the Orion OB1 association, the most systematic search was longer associated with their parent molecular clouds, mak- that done with the 1-m Kiso Schmidt (e.g., Wiramihardja ing it difficult to sort them out from the field population. et al., 1989, 1993), covering roughly 150 deg2 and detecting Therefore, a particular combination of various instruments ~1200 emission line stars, many of which were argued to be and techniques is required to reliably single out the low- likely TTS. Weaver and Babcock (2004) recently identified mass PMS stars. The main strategies that have been used to 63 Hα-emitting objects in a deep objective prism survey of identify these populations are objective prism surveys, X- the σ Orionis region. ray emission, proper motions, and, more recently, variability The main limitation of this technique is the strong bias surveys. toward Hα-strong PMS stars; few WTTS can be detected at the resolution of objective prisms (cf. Briceño et al., 2.1. Objective Prism Surveys 1999). Briceño et al. (2001) find that only 38% of the 151 Kiso Hα sources falling within their ~34 deg2 survey area The TTS originally were identified as stars of late spec- in the Orion OB1a and 1b subassociations are located above tral types (G-M), with strong emission lines (especially Hα) the ZAMS in color-magnitude diagrams, and argue that the and erratic light variations, spatially associated with regions Kiso survey is strongly contaminated by foreground main- of dark nebulosity (Joy, 1945). Stars resembling the original sequence stars (largely dMe stars). The spatial distribution variables first identified as TTS are currently called “strong of the Kiso sources has been useful to outline the youngest emission” or Classical TTS (CTTS). Subsequent spectro- regions in Orion, where the highest concentrations of CTTS scopic studies of the Ca II H and K lines and the first X-ray are located (Gómez and Lada, 1998), but these samples can observations with the Einstein X-ray observatory (Feigelson be dominated by field stars in regions far from the molecu- and De Campli, 1981; Walter and Kuhi, 1981) revealed sur- lar clouds, in which the CTTS/WTTS fraction is small. prisingly strong X-ray activity in TTS, exceeding the solar Therefore, as with other survey techniques, objective prism levels by several orders of magnitude, and also revealed a studies require follow-up spectroscopy to confirm member- population of X-ray strong objects lacking the optical sign- ship. posts of CTTS, like strong Hα emission. These stars, initially called “naked-T Tauri stars” (Walter and Myers, 1986), 2.2. X-Ray Surveys are now widely known as “weak-line” TTS after Herbig and Bell (1988). The CTTS/WTTS dividing line was set Young stars in all evolutionary stages, from class I pro- at W(Hα) = 10 Å. In general, the excess Hα emission in tostars to ZAMS stars, show strong X-ray activity [for re- WTTS seems to originate in enhanced solar-type magnetic cent reviews on the X-ray properties of YSOs see Feigelson activity (Walter et al., 1988), while the extreme levels ob- and Montmerle (1999) and Favata and Micela (2003)]. served in CTTS can be explained by a combination of en- After the initial Einstein studies, the ROSAT and ASCA hanced chromospheric activity and emission coming from X-ray observatories increased considerably the number of accretion shocks in which material from a circumstellar disk observed star-forming regions, and thereby the number of is funneled along magnetic field lines onto the stellar photo- known X-ray-emitting TTS. Today, XMM-Newton and sphere (section 5). Recently, White and Basri (2003) revis- Chandra allow X-ray studies of star-forming regions at un- ited the WTTS/CTTS classification and suggested a modi- precedented sensitivity and spatial resolution. fied criterion that takes into account the contrast effect in X-ray observations are a well-established tool to find Hα emission as a function of spectral type in stars cooler young stars. For nearby OB associations, which typically than late K. cover areas in the sky much larger than the field-of-view The strong Hα emission characteristic of low-mass young of X-ray observatories, deep and spatially complete obser- stars, and in particular of CTTS, encouraged early large-scale vations are usually not feasible. However, large-scale shal- searches using photographic plates and objective prisms on low surveys have been conducted with great success. The widefield instruments like Schmidt telescopes [e.g., San- ROSAT All Sky Survey (RASS) provided coverage of the duleak (1971), in Orion]. These very-low-resolution spec- whole sky in the 0.1–2.4-keV soft X-ray band. With a mean troscopic surveys [typical dispersions of ~500 Å/mm to limiting flux of about 2 × 10–13 erg s–1 cm2 this survey pro- ~1700 Å/mm at Hα (cf. Wilking et al., 1987; Briceño et al., vided a spatially complete, flux-limited sample of X-ray 1993)] provided large area coverage, allowed estimates of sources that led to the detection of hundreds of candidate 348 Protostars and Planets V

PMS stars in star-forming regions all over the sky (see Neu- candidates would be bona fide association members, and häuser, 1997). indeed Mamajek et al. (2002) found that 22/30 (73%) of a The X-ray luminosities of young stars for a given age, subsample of candidates, located in Sco-Cen, could be spec- mass, and bolometric luminosity can differ by several orders troscopically confirmed as PMS stars. Hoogerwerf (2000) of magnitude. Until recently it was not even clear whether used the ACT and TRC proper motion catalogs (p.m. er- all young stars are highly X-ray active, or whether an “X- rors = 3 mas yr–1) to identify thousands of candidate Sco- ray quiet” population of stars with suppressed magnetic ac- Cen members down to V ~ 12. Unfortunately, the vast ma- tivity may exist, which would have introduced a serious bias jority of stars in the ACT and TRC catalogs, and their de- in any X-ray selected sample. The Chandra Orion Ultradeep scendant (Tycho-2), do not have known spectral types or Project (Getman et al., 2005), a 10-day-long observation parallaxes (in contrast to the Hipparcos catalog), and hence of the Orion Nebular Cluster, has provided the most com- the contamination level is large. Mamajek et al. (2002) con- prehensive dataset ever acquired on the X-ray emission of ducted a spectroscopic survey of an X-ray and color-mag- PMS, and solved this question by providing definitive in- nitude-selected subsample of the Hoogerwerf proper-mo- formation on the distribution of X-ray luminosities in young tion-selected sample and found that 93% of the candidates stars. It found no indications for “X-ray quiet” TTS, and were bona fide PMS association members. The high-qual- ≥ established that 50% of the TTS have log (LX/Lbol) –3.5, ity proper motions also enabled the estimate of individual ≥ while 90% have log (LX/Lbol) –4.5 (Preibisch et al., 2005; parallaxes to the Upper Centaurus Lupus (UCL) and Lower see also chapter by Feigelson et al.). Since the RASS flux Centarus Crux (LCC) members, reducing the scatter in the limit corresponds to X-ray luminosities of about 5 × 1029 erg HR diagram (Mamajek et al., 2002). In a survey of 115 can- s–1 at the distance of the nearest OB associations (~140 pc), didate Upper Sco members selected solely via STARNET this implies that the RASS data are essentially complete only proper motions (p.m. errors of ~5 mas yr–1), Preibisch et al. for M = 1 M PMS stars in those regions, while only a frac- (1998) found that none were PMS stars. The lesson learned tion of the X-ray brightest subsolar-mass PMS stars are de- appears to be that proper motions alone are insufficient for tected. A caveat of the RASS surveys for PMS stars is that efficiently identifying low-mass members of nearby OB these samples can be significantly contaminated by fore- associations. However, when proper motions are used in ground, X-ray active zero-age main-sequence stars (Briceño conjunction with color-magnitude, X-ray, spectral type, or et al., 1997). These limitations have to be kept in mind when parallax data (or some combination thereof), finding low- working with X-ray selected samples; at any rate, follow- mass associations can be a very efficient task. up observations are necessary to determine the nature of the objects. 2.4. Photometric Surveys: Single-Epoch Observations 2.3. Proper Motion Surveys Single-epoch photometric surveys are frequently used to The recent availability of ever-deeper, all-sky catalogs select candidate low-mass members of young clusters or of proper motions (like Hipparcos and the Tycho family of associations. Most studies use broadband, optical filters that catalogs) has aided the effort in identifying the low-mass are sensitive to the temperatures of G-, K-, and M-type stars. members of the nearest OB associations. The proper mo- Near-IR color-magnitude diagrams (CMDs) are not as use- tions of members of a few of the nearest OB associations ful for selecting low-mass PMS stars because NIR colors are on the order of tens of mas yr–1 (de Zeeuw et al., 1999). are similar for all late-type stars. With proper motions whose errors are less than a few mas Candidate low-mass association members are usually se- yr –1, one can attempt to kinematically select low-mass mem- lected by their location in the CMD above the zero-age main bers of nearby associations. The nearest OB association, sequence (ZAMS). This locus is usually defined by either Sco-Cen, has been the most fruitful hunting ground for iden- a known (spectroscopically confirmed) population of PMS tifying low-mass members by virtue of their proper motions. stars or because the PMS population of the association is Current proper motion catalogs (e.g., Tycho-2, UCAC) clearly visible as a concentration on the CMD (e.g., Fig. 1). are probably adequate to consider kinematic selection of Single-epoch photometry is most effective in regions such low-mass stars in at least a few other nearby groups (e.g., as σ Ori or the ONC where the proximity and youth of the Vel OB2, Tr 10, α Per, Cas-Tau, Cep OB6). The very small cluster make the low-mass PMS members brighter than the (<10 mas yr–1) proper motions for some of the other nearby bulk of the field stars at the colors of K- and M-type stars. OB associations (e.g., Ori OB1, Lac OB1, Col 121) will pre- The main advantage of photometric selection is that for clude any attempts at efficient selection of low-mass mem- a specified amount of time on any given telescope a region bers via proper motions, at least with contemporary astro- of the sky can be surveyed to a fainter limit than can be metric catalogs. done by a variability survey or a spectroscopic survey. Also, The Hipparcos survey of the nearest OB associations by photometric selection can identify low-mass association de Zeeuw et al. (1999) was able to identify dozens of FGK- members with very low amplitude variability. The disadvan- type stars as candidate members. de Zeeuw et al. (1999) tage of single-epoch photometric selection is that there is predicted that ~37/52 (71%) of their GK-type Hipparcos inevitably some contamination by foreground field stars and Briceño et al.: Low-mass Populations in OB Associations 349

over an area of >150 deg2 in the Orion OB1 association. (a) (b) In their first release, based on some 25 epochs and spanning an area of 34 deg2, they identified ~200 new, spectroscopically confirmed, low-mass members, and a new 10- m.y.-old clustering of stars around the star 25 Ori (Briceño et al., in preparation). McGehee et al. (2005) analyzed nine repeated observations over 25 deg2 in Orion, obtained with the Sloan Digital Sky Survey (SDSS). They selected 507 stars that met their variability criterion in the SDSS g-band. They did not obtain follow-up spectra of their candidates, rather, they apply their observations in a statistical sense to search for photometric accretion-related signatures in their lower-mass candidate members. Slesnick et al. (2005) are using the Quest II CCD Mosaic Camera on the Palomar Schmidt to conduct a BRI, multi-epoch survey of ~200 deg2 in Upper Sco, and Carpenter et al. (2001) has used repeated Fig. 1. (a) V vs. V-IC color-magnitude diagram of 9556 stars in observations made with 2 MASS in a 0.86° × 6° strip cen- 0.89 deg2 around σ Ori (from Sherry et al., 2004). The solid line tered on the ONC to study the near-IR variability of a large is a 2.5-m.y. isochrone (Baraffe et al., 1998, 2001) at a distance sample of young, low-mass stars. of 440 pc (de Zeeuw et al., 1999), extending from 1.2 M at V ~ With more wide-angle detectors on small- and medium- 13.5 down to ~0.2 M at V ~ 18 (the completeness limit, indicated sized telescopes, and projects like LSST coming on line by the dashed line). This isochrone marks the expected position of within less than a decade, variability promises to grow as the PMS locus for Orion OB1b. There is a clear increase in the an efficient means of selecting large samples of candidate density of stars around the expected position of the PMS locus. (b) The same color-magnitude diagram (CMD) for the 0.27 deg2 low-mass PMS stars down to much lower masses than avail- control fields from Sherry et al. (2004). The isochrone (solid line) able to all-sky X-ray surveys like ROSAT, and without the is the same as in the left panel. The dashed line marks the fainter bias toward CTTS of objective prism studies. However, as completeness limit of the control fields. with every technique there are limitations involved, e.g., temporal sampling and a bias toward variables with larger amplitudes, especially at the faint end, are issues that need to be explored. background giants. As with the other techniques, it is impos- sible to securely identify any individual star as a low-mass 2.6. Spectroscopy member of the association without spectroscopic follow-up. In small areas with a high density of low-mass association As already emphasized in the previous paragraphs, low- members such as the σ Ori cluster or the ONC, single-epoch to moderate-resolution follow-up spectroscopy is essential photometry can effectively select the low-mass population to confirm membership of PMS stars. However, the obser- because field star contamination is fairly small (Sherry et vational effort to identify the widespread population of PMS al., 2004; Kenyon et al., 2005). But in large areas with a stars among the many thousands of field stars in the large lower density of low-mass association members, such as areas spanned by nearby OB associations is huge, and up Orion OB1b (Orion’s belt) or Orion OB1a (northwest of the to recently has largely precluded further investigations. With belt), the field star contamination can be large enough to the advent of extremely powerful multiple object spectro- make it difficult to even see the PMS locus. graphs, such as 2dF at the Anglo-Australian Telescope, Hy- dra on the WIYN 3.5-m and the CTIO 4-m telescopes, and 2.5. Photometric Surveys: Variability now Hectospec on the 6.5-m MMT, large-scale spectroscopic surveys have now become feasible. Variability in T Tauri stars has been intensively studied One of the most powerful approaches to unbiased sur- over the years (e.g., Herbst et al., 1994), but mostly as fol- veys of young low-mass stars is to use the presence of low-up observations of individual young stars that had been strong 6708 Å Li I absorption lines as a diagnostic of the identified by some other means. Building on the availabil- PMS nature of a candidate object (e.g., Dolan and Mathieu, ity of large-format CCD cameras installed on widefield tele- 1999, 2001). Because Li I is strongly diminished in very scopes it has now become feasible to conduct multi-epoch, early phases of stellar evolution, a high Li content is a re- photometric surveys that use variability to pick out candi- liable indication for the youth of a star (e.g., Herbig, 1962; date TTS over the extended areas spanned by nearby OB D’Antona and Mazzitelli, 1994). However, Li depletion is associations. In Orion, two major studies have been con- not only a function of stellar age, but also of stellar mass ducted over the past few years. Briceño et al. (2001, 2005a) and presumably even depends on additional factors like have done a VRI variability survey using the Quest I CCD stellar rotation (cf. Soderblom, 1996). Not only PMS stars, Mosaic Camera installed on the Venezuela 1-m Schmidt, but also somewhat older, although still relatively young 350 Protostars and Planets V zero-age main-sequence stars, e.g., the G- and K-type stars in the ~108-yr-old Pleiades (cf. Soderblom et al., 1993), can display Li absorption lines. In order to classify stars as PMS, one thus has to consider a spectral type dependent threshold for the Li line width. Such a threshold can be defined by the upper envelope of Li measurements in young clusters of main-sequence stars with ages between ~30 m.y. and a few 100 m.y. such as IC 2602, IC 4665, IC 2391, α Per, and the Pleiades. Any star with a Li line width considerably above this threshold should be younger than ~30 m.y. and can therefore be classified as a PMS object. In addition to Li, other spectroscopic signatures can be used as youth indicators, such as the K I and Na I absorption lines that are typically weaker in low-mass PMS objects compared to field M-type dwarfs (Martín et al., 1996; Luh- man, 1999), and radial velocities (if high-resolution spectra are available).

3. AGES AND AGE SPREADS OF LOW-MASS STARS

Stellar ages are usually inferred from the positions of the Fig. 2. HR diagram for the Upper Sco association members from stars in the HR diagram by comparison with theoretical the study of Preibisch et al. (2002). The lines show the evolu- PMS evolutionary models. Photometry of young clusters tionary tracks from the Palla and Stahler (1999) PMS models, and associations has shown that low-mass members occupy some labeled by their masses in solar units. The thick solid line a wide swath on the CMD or the H-R diagram (NGC 2264, shows the main sequence. The 5-m.y. isochrone is shown as the ONC, σ Ori). A common interpretation has been that this dashed line; it was composed from the high-mass isochrone from Bertelli et al. (1994) for masses 6–30 M , the Palla and Stahler spread is evidence of an age spread among cluster mem- (1999) PMS models for masses 1–6 M , and the Baraffe et al. bers (e.g., Palla and Stahler, 2000). This interpretation as- (1998) PMS models for masses 0.02–1 M . The gray shaded band sumes that the single-epoch observed colors and magnitudes shows the region in which one expects 90% of the member stars of low-mass PMS stars accurately correspond to the tem- to lie, based on the assumption of a common age of 5 m.y. for all peratures and luminosities of each star. However, it has to stars and taking proper account of the uncertainties and the effects be strongly emphasized that the masses and especially the of unresolved binaries (for details see text). ages of the individual stars read off from their position in these diagrams are generally not identical to their true masses and ages. Several factors can cause considerable deviations of an individual star’s position in the HR-dia- 5-m.y. isochrone. Not only the majority of the low-mass gram from the locations predicted by theoretical models for stars, but also most of the intermediate- and high-mass stars, a given age and mass. Low-mass PMS stars exhibit vari- lie close to or on the 5-m.y. isochrone. There clearly is a ability ranging from a few tenths in the WTTS up to sev- considerable scatter that may seem to suggest a spread in eral magnitudes in the CTTS (Herbst et al., 1994). Further- stellar ages. In the particular case of Upper Sco shown here, more, binarity can add a factor of 2 (~0.75 mag) spread to in addition to the other effects mentioned above, the most the CMD. For regions spanning large areas on the sky there important factor for the apparent scatter is the relatively will be an additional spread caused by the distribution of large spread of individual stellar distances [~±20 pc around distances along the line of sight among individual stars. the mean value of 145 pc (de Bruijne, 1999, and personal Another source of uncertainty are the calibrations used to communication)] in this very nearby and extended region, derive bolometric luminosities and effective temperatures, which causes the luminosities to be either over- or under- and finally, the choice of evolutionary tracks used, that in estimated when a single distance is adopted for all sources. some cases yield different ages for intermediate (~2–5 M ) A detailed discussion and statistical modeling of these ef- and low-mass stars located in the same region [e.g., the ONC fects is given in Preibisch and Zinnecker (1999) and Prei- (Hillenbrand, 1997; see also Hartmann, 2001)]. Therefore, bisch et al. (2002). In the later work these authors showed any age estimates in young star-forming regions must ac- that the observed HR diagram for the low-mass stars in count for the significant spread that a coeval population will Upper Sco is consistent with the assumption of a common necessarily have on the CMD, which translates into a spread stellar age of about 5 m.y.; there is no evidence for an age on the HR diagram. dispersion, although small ages spreads of ~1–2 m.y. cannot As an example, Fig. 2 shows the HR diagram containing be excluded by the data. Preibisch et al. (2002) showed that all Upper Sco association members from Preibisch et al. the derived age is also robust when taking into account the (2002); the diagram also shows the main sequence and a uncertainties of the theoretical PMS models. It is remarkable Briceño et al.: Low-mass Populations in OB Associations 351

that the isochronal age derived for the low-mass stars is consistent with previous and independent age determinations based on the nuclear and kinematic ages of the massive stars (de Zeeuw and Brand, 1985; de Geus et al., 1989), which also yielded 5 m.y. This very good agreement of the independent age determinations for the high-mass and the low- mass stellar population shows that low- and high-mass stars are coeval and thus have formed together. Furthermore, the absence of a significant age dispersion implies that all stars in the association have formed more or less simultaneously. Therefore, the star-formation process must have started rather suddenly and everywhere at the same time in the association, and also must have ended after at most a few million years. The star-formation process in Upper Sco can thus be considered as a burst of star formation. Sherry (2003) compared the observed spread across the σ V vs. V-IC CMD of low-mass members of the Ori cluster to the spread that would be expected based upon the known variability of WTTS, the field binary fraction, the photometric errors of his survey, and a range of simple star-formation histories. The observed spread was consistent with the predicted spread of an isochronal population. Sherry concluded Fig. 3. Distribution of ages of ~1000 newly identified TTS in a that the bulk of the low-mass stars must have formed over ~60 deg2 area spanning the Orion OB1a and 1b subassociations a period of less than ~1 m.y. A population with a larger age (Briceño et al., in preparation). Ages were derived using the spread would have been distributed over a larger region of Baraffe et al. (1998) and Siess et al. (2000) evolutionary tracks the CMD than was observed. and isochrones. The mean ages are ~4 m.y. for Ori OB1b and Burningham et al. (2005) also examined the possibility ~8 m.y. for Ori OB1a. of an age spread among members of the σ Ori cluster. They used two epoch R and i’ observations of cluster members taken in 1999 and 2003 to estimate the variability of each distributions of high-mass and solar-type stars in the region cluster member. They then constructed a series of simple show several critical features: (1) both high- and low-mass models with a varying fraction of equal mass binaries. They star formation began concurrently in the center of the SFR found that the observed spread on the CMD was too large to roughly 6–8 m.y. ago; (2) low-mass star formation ended in be fully accounted for by the combined effects of observa- the vicinity of λ Ori roughly 1 m.y. ago; (3) low-mass star tional errors, variability (over 1–4 years), and binaries. They formation rates near the B30 and B35 clouds reached their conclude that the larger spread on the CMD could be ac- maxima later than did low-mass star formation in the vicin- counted for by either a longer period of accretion driven ity of λ Ori; (4) low-mass star formation continues today variability, an age spread of <2 m.y. (using a distance of near the B30 and B35 clouds. As with the Ori OB1 asso- 440 pc or 4 m.y. using a distance of 350 pc), or a combina- ciations, this varied star formation history reflects the rich tion of long-term variability and a smaller age spread. How- interplay of the massive stars and the gas. ever, their result is not necessarily in contradiction with the The accumulated evidence for little, if any, age spreads findings of Sherry (2003), especially if the actual variability in various star-forming regions provides a natural explana- is larger than they estimate based on their two epoch obser- tion for the “post-T Tauri problem” [the absence of “older” vations. TTS in star-forming regions like Taurus, assumed to have In the Orion OB1 association, the V and IC-band CMDs been forming stars for up to tens of millions of years (Her- for spectroscopically confirmed TTS (Briceño et al., 2005a; big, 1978)], and at the same time implies relatively short Briceño et al., in preparation), which mitigate the spread lifetimes for molecular clouds. These observational argu- caused by the variability of individual sources by plotting ments support the evolving picture of star formation as a their mean magnitudes and colors, provide evidence that the fast and remarkably synchronized process in molecular Ori OB1b subassociation is older (~4 m.y.) than the often clouds (section 6). quoted age of about 1.7 m.y. (Brown et al., 1994), with a rather narrow age distribution regardless of the tracks used 4. THE INITIAL MASS FUNCTION (Fig. 3). The same dataset for the older OB1a subassociation IN OB ASSOCIATIONS also shows a narrow range of ages, with a mean value of ~8 m.y. The IMF is the utmost challenge for any theory of star The λ Ori region also shows that the age distribution and formation. Some theories suggest that the IMF should vary star-formation history is spatially dependent. Dolan and systematically with the star-formation environment (Larson, Mathieu (2001) (see also Lee et al., 2005) found that age 1985), and for many years star formation was supposed to 352 Protostars and Planets V be a bimodal process (e.g., Shu and Lizano, 1988) according to which high- and low-mass stars should form in to- tally different sites. For example, it was suggested that increased heating due to the strong radiation from massive stars raises the Jeans mass, so that the bottom of the IMF would be truncated in regions of high-mass star formation. In contrast to the rather quiescent environment in small low- mass clusters and T associations (like the Taurus molecular clouds), forming stars in OB association are exposed to the strong winds and intense UV radiation of the massive stars, and, after a few million years, also affected by supernova explosions. In such an environment, it may be harder to form low-mass stars, because, e.g., the lower-mass cloud cores may be completely dispersed before protostars can even begin to form. Although it has been long established that low-mass stars can form alongside their high-mass siblings (e.g., Herbig, 1962) in nearby OB associations, until recently it was not well known what quantities of low-mass stars are produced in OB environments. If the IMF in OB associations is not truncated and similar to the field IMF, it would follow that most of their total stellar mass (>60%) is found in low-mass (<2 M ) stars. This would then imply that most of the cur- Fig. 4. Comparison of the mass function derived for the Upper rent galactic star formation is taking place in OB associa- Sco association with different mass function measurements for the tions. Therefore, the typical environment for forming stars field (from Preibisch et al., 2002). The Upper Sco mass function (and planets) would be close to massive stars and not in iso- is shown three times by the solid dots connected with the dotted lated regions like Taurus. lines, multiplied by arbitrary factors. The middle curve shows the In Fig. 4 we show the empirical mass function for Upper original mass function; the solid line is our multipart power law fit. The upper curve shows our mass function multiplied by a factor Sco as derived in Preibisch et al. (2002), for a total sample of 30 and compared to the Scalo (1998) IMF (solid line); the gray of 364 stars covering the mass range from 0.1 M up to shaded area delimited by the dashed lines represents the range 20 M . The best-fit multipart power law function for the allowed by the errors of the model. The lower curve shows our probability density distribution is given by mass function multiplied by a factor of 1/30 and compared to the Kroupa (2002) IMF (solid line).

M–0.9 ± 0.2 for 0.1 ≤ M/M < 0.6 dN ∝ M–2.8 ± 0.5 for 0.6 ≤ M/M < 2 (1) dM M–2.6 ± 0.5 for 0.2 ≤ M/M < 20 Barrado y Navascues et al. (2004) combined their deep im- aging data with the surveys of Dolan and Mathieu (1999, 2001) (limited to the same area) to obtain an initial mass or, in shorter notation, α (0.1–0.6) = –0.9 ± 0.2, α (0.6– function from 0.02 to 1.2 M . They find that the data indi- 2.0) = –2.8 ± 0.5, α (2.0–20) = –2.6 ± 0.3. For comparison, cate a power law index of α = –0.60 ± 0.06 across the stel- the plot also shows two different field IMF models, the lar-substellar limit and a slightly steeper index of α = 0.86 ± Scalo (1998) model, which is given by α (0.1–1) = –1.2 ± 0.05 over the larger mass range of 0.024 M to 0.86 M , 0.3, α (1–10) = –2.7 ± 0.5, α (10–100) = –2.3 ± 0.5, and much as is found in other young regions. the Kroupa (2002) model with α (0.02–0.08) = –0.3 ± 0.7, Over the entire λ Ori star-forming region, Dolan and α (0.08–0.5) = –1.3 ± 0.5, α (0.5–100) = –2.3 ± 0.3. Mathieu (2001) were able to clearly show that the IMF has While the slopes of the fit to the empirical mass function a spatial dependence. Dolan and Mathieu (1999) had found of Upper Sco are not identical to those of these models, they that within the central ~3.5 pc around λ Ori the low-mass are well within the ranges of slopes derived for similar mass stars were deficient by a factor of 2 compared to the field ranges in other young clusters or associations, as compiled in IMF. Outside this central field, Dolan and Mathieu (2001) Kroupa (2002). Therefore, it can be concluded that, within showed the low-mass stars to be overrepresented compared the uncertainties, the general shape of the Upper Sco mass to the Miller and Scalo (1978) IMF by a factor of 3. A simi- function is consistent with recent field star and cluster IMF lar overrepresentation of low-mass stars is also found at sig- determinations. nificant confidence levels when considering only stars as- In the Orion OB1 association Sherry et al. (2004) find sociated with B30 and B35. that the mass function for the σ Ori cluster is consistent with Thus the global IMF of the λ Ori SFR resembles the field, the Kroupa (2002) IMF. In the immediate vicinity of λ Ori, while the local IMF appears to vary substantially across the Briceño et al.: Low-mass Populations in OB Associations 353

region. No one place in the λ Ori SFR creates the field IMF falling material also produces the observed broadened and by itself. Only the integration of the star-formation process P Cygni profiles observed in hydrogen lines (Muzerolle et over the entire region produces the field IMF. al., 1998a,b, 2001). Disk accretion rates for most CTTS are on the order of 10–8 M yr–1 at ages of 1–2 m.y. (e.g., Gull- 5. DISK EVOLUTION bring et al., 1998; Hartmann, 1998; Johns-Krull and Val- enti, 2001). The presence of circumstellar disks around low-mass Comparative studies of near-IR emission and accretion- pre-main-sequence stars appears to be a natural consequence related indicators (Hα and Ca II emission, UV excess emis- of the star-formation process; these disks play an important sion) at ages ~1–10 m.y. offer insight into how the inner- role both in determining the final mass of the star and as most part of the disk evolves. One way to derive the fraction the potential sites for planet formation. Although we have of stellar systems with inner disks is counting the number a good general understanding of the overall processes in- of objects showing excess emission in the JHKL near-IR volved, many important gaps still remain. For instance, the bands. The availability of 2MASS has made JHK studies timescales for mass accretion and disk dissipation are still of young populations over wide spatial scales feasible. More matters of debate. An example is the discovery of a seem- recently, the emission of T Tauri stars in the Spitzer IRAC ingly long-lived accreting disk around the ~25-m.y.-old, and MIPS bands has been characterized by Allen et al. late-type (M3) star St 34 located in the general area of the (2004) and Hartmann et al. (2005b). Another approach is Taurus dark clouds (Hartmann et al., 2005a; White and Hil- determining the number of objects that exhibit strong Hα lenbrand, 2005). and Ca II emission (CTTS), or UV excesses; these figures How circumstellar disks evolve and whether their evo- provide an indication of how many systems are actively ac- lution is affected by environmental conditions are questions creting from their disks. that at present can only be investigated by looking to low- So far, the most extensive studies of how disk fractions mass PMS stars, and in particular the best samples can now change with time have been conducted in the Orion OB1 be drawn from the various nearby OB associations. First, association. Hillenbrand et al. (1998) used the Ic-K color low-mass young stars like T Tauri stars, in particular those to derive a disk fraction of 61–88% in the ONC, and the with masses ~1 M , constitute good analogs of what the Ca II lines in emission, or “filled-in,” as a proxy for deter- conditions may have been in the early solar system. Sec- mining an accretion disk frequency of ~70%. Rebull et al. ond, as we have discussed in section 3, OB associations can (2000) studied a region on both sides of the ONC and de- harbor many stellar aggregates, with distinct ages, such as in termined a disk accretion fraction in excess of 40%. Lada Orion OB1, possibly the result of star-forming events (trig- et al. (2000) used the JHKL bands to derive a ONC disk gered or not) occurring at various times throughout the orig- fraction of 80–85% in the low-mass PMS population. Lada inal GMC. Some events will have produced dense clusters et al. (2004) extended this study to the substellar candidate while others are responsible for the more spread out popula- members and found a disk fraction of ~50%. In their on- tion. The most recent events are easily recognizable by the going large-scale study of the Orion OB1 association, Cal- very young (<1 m.y.) stars still embedded in their natal gas, vet et al. (2005a) combined UV, optical, JHKL, and 10-µm while older ones may be traced by the ~10-m.y. stars that measurements in a sample of confirmed members of the 1a have long dissipated their parent clouds. This is why these and 1b subassociations to study dust emission and disk ac- regions provide large numbers of PMS stars in different envi- cretion. They showed evidence for an overall decrease in ronments, but presumably sharing the same “genetic pool,” IR emission with age, interpreted as a sign of dust evolution that can allow us to build a differential picture of how disks between the disks in Ori OB1b (age ~4 m.y.), Ori OB1a evolve from one stage to the next. (age ~8 m.y.), and those of younger populations like Tau- Disks are related to many of the photometric and spectro- rus (age ~2 m.y.). Briceño et al. (2005b) used IRAC and scopic features observed in T Tauri stars. The IR emission MIPS on Spitzer to look for dusty disks in Ori OB1a and originates by the contribution from warm dust in the disk, 1b. They confirm a decline in IR emission by the age of heated at a range of temperatures by irradiation from the Orion OB1b, and find a number of “transition” disk systems star and viscous dissipation (e.g., Meyer et al., 1997). The (~14% in 1b and ~6% in 1a), objects with essentially photo- UV excesses, excess continuum emission (veiling), irregular spheric fluxes at wavelengths ≤4.5 µm and excess emission photometric variability, and broadened spectral line profiles at longer wavelengths. These systems are interpreted as (particularly in the hydrogen lines and others like Ca II) are showing signatures of inner disk clearing, with optically thin explained as different manifestations of gas accretion from inner regions stretching out to one or a few AU (Calvet et a circumstellar disk. In the standard magnetospheric model al., 2002, 2005b; Uchida et al., 2004; D’Alessio et al., (Königl, 1991), the accretion disk is truncated at a few stel- 2005); the fraction of these transition disks that are still lar radii by the magnetic field of the star; the disk mate- accreting is low (~5–10%), hinting at a rapid shutoff of the rial falls onto the photosphere along magnetic field lines at accretion phase in these systems [similar results have been supersonic velocities, creating an accretion shock that is obtained by Sicilia-Aguilar et al. (2006) in Cep OB2; see thought to be largely responsible for the excess UV and below]. Haisch et al. (2000) derived an IR-excess fraction continuum emission (Calvet and Gullbring, 1998). The in- of ~86% in the <1-m.y.-old NGC 2024 embedded cluster. 354 Protostars and Planets V

Their findings indicate that the majority of the sources that formed in NGC 2024 are presently surrounded by, and were likely formed with, circumstellar disks. One of the more surprising findings in the λ Ori region was that, despite the discovery of 72 low-mass PMS stars within 0.5° (~3.5 pc) of λ Ori, only two of them showed strong Hα emission indicative of accretion disks (Dolan and Mathieu, 1999). Dolan and Mathieu (2001) expanded on this result by examining the distribution of Hα emission along an axis from B35 through λ Ori to B30. The paucity of Hα emission-line stars continues from λ Ori out to the two dark clouds, at which point the surface density of Hα emission-line stars increases dramatically. Even so, many of the Hα stars associated with B30 and B35 have ages similar to PMS stars found in the cluster near λ Ori. Yet almost none of the latter show Hα emission. This strongly suggests that the absence of Hα emission from the central PMS stars is the result of an environmental influence linked to the luminous OB stars. The nearby Scorpius-Centaurus OB association [d ~ Fig. 5. Inner disk fraction around low-mass young stars (0.2 < 130 pc (de Zeeuw et al., 1999)] has been a natural place to M < 1 M ) as a function of age in various nearby OB associa- search for circumstellar disks, especially at somewhat older tions (solid dots). The ONC point is an average from estimates by ages (up to ~10 m.y.). Moneti et al. (1999) detected excess Hillenbrand et al. (1998), Lada et al. (2000), and Haisch et al. emission at the ISOCAM 6.7- and 15-µm bands in 10 X-ray (2001). Data in the JHKL bands from Haisch et al. (2001) were selected WTTS belonging to the more widely spread popu- used for IC 348, NGC 2264, and NGC 2362. Note that for the more distant regions like NGC 2362 and NGC 2264, the mass lation of Sco-Cen, albeit at levels significantly lower than completeness limit is ~1 M . Results for Orion OB1a and 1b are in the much younger (age ~1 m.y.) Chamaeleon I T associa- from JHK data in Briceño et al. (2005a). In Tr 37 and NGC 7160 tion. In a sample of X-ray and proper motion-selected late- we used values by Sicilia-Aguilar et al. (2005, 2006) from accre- type stars in the Lower Centaurus Crux (LCC, age ~17 m.y.) tion indicators and JHK (3.6 µm) measurements. As a compari- and Upper Centaurus Lupus (UCL, age ~15 m.y.) in Sco- son, we also plot disk fractions (open squares) derived in Taurus Cen, Mamajek et al. (2002) find that only 1 out of 110 PMS (from data in Kenyon and Hartmann, 1995) and for Chamaeleon solar-type stars shows both enhanced Hα emission and a from Gómez and Kenyon (2001). A conservative error bar is indi- K-band excess indicative of active accretion from a truncated cated at the lower left corner. circumstellar disk, suggesting timescales of ~10 m.y. for halting most of the disk accretion. Chen et al. (2005) obtained Spitzer Space Telescope MIPS observations of 40 F- in optically thick disks as a function of age, with faster de- and G-type common proper motion members of the Sco- creases in the IR emission at shorter wavelengths, suggestive Cen OB association with ages between 5 and 20 m.y. They of a more rapid evolution in the inner disks. detected 24 µm excess emission in 14 objects, corresponding All these studies agree on a general trend toward rapid to a disk fraction of ≥35%. inner disk evolution (Fig. 5); it looks like the disappearance In Perseus, Haisch et al. (2001) used JHKL observations of the dust in the innermost parts of the disk, either due to to estimate a disk fraction of ~65% in the 2–3-m.y.-old grain growth and settling or to photoevaporation (Clarke (Luhman et al., 1998) IC 348 cluster, a value lower than in et al., 2001), is followed by a rapid shutoff of accretion. the younger NGC 2024 and ONC clusters in Orion, sug- However, both these investigations and newer findings also gestive of a timescale of 2–3 m.y. for the disappearance of reveal exceptions that lead to a more complex picture. Mu- approximately one-third of the inner disks in IC 348. The zerolle et al. (2005) find that 5–10% of all disk sources in ~4-m.y.-old Tr 37 and the ~10-m.y.-old NGC 7160 clus- the ~1-m.y.-old NGC 2068 and 2071 clusters in Orion show ters in Cep OB2 have been studied by Sicilia-Aguilar et al. evidence of significant grain growth, suggesting a wide (2005), who derive an accreting fraction of ~40% for Tr 37, variation in timescales for the onset of primordial disk evo- and 2–5% (one object) for NGC 7160. Sicilia-Aguilar et al. lution and dissipation. This could also be related to the ex- (2006) used Spitzer IRAC and MIPS observations to fur- istence of CTTS and WTTS in many regions at very young ther investigate disk properties in Cep OB2. About 48% of ages, the WTTS showing no signatures of inner, optically the members exhibit IR excesses in the IRAC bands, con- thick disks. Did some particular initial conditions favor very sistent with their inferred accreting disk fraction. They also fast disk evolution in these WTTS (binarity)? Finally, the find a number of “transition” objects (10%) in Tr 37. They detection of long-lived disks implies that dust dissipation interpret their results as evidence for differential evolution and the halting of accretion do not necessarily follow a uni- Briceño et al.: Low-mass Populations in OB Associations 355

versal trend. This may have implications for the formation selves can form very quickly from atomic gas, in 1–2 m.y. of planetary systems; if slow accretion processes are the (see chapter by Ballesteros-Paredes et al.)]. The unbound dominant formation mechanism for jovian planets then long- state of the GMC ensures that the whole region is dispers- lived disks may be ideal sites to search for evidence for pro- ing while it is forming stars or star clusters in locally com- toplanets. pressed sheets and filaments, due to the compressive nature of supersonic turbulence (“converging flows”). The 6. THE ORIGIN OF OB ASSOCIATIONS: mass fraction of compressed cloud gas is low, affecting only BOUND VS. UNBOUND STELLAR GROUPS ~10% of the GMC. In this model, the spacing of the OB stars that define an OB association would be initially large Blauuw (1964) in his masterly review first provided some (larger than a few parsec), rather than compact (<1 pc) as clues as to the origin of OB associations, alluding to two in the KLL model. Furthermore, no sequential star forma- ideas: “originally small, compact bodies with dimensions tion triggered by the thermal pressure of expanding HII of several parsec or less” or “regional star formation, more regions “a la Elmegreen and Lada (1977)” may be needed or less simultaneous, but scattered over different parts of a to create OB subgroups, unlike in the KLL model. Delayed large cloud complex.” Similarly, there are at present two SN triggered star formation in adjacent transient clouds at (competing) models that attempt to explain the origin of OB some distance may still occur and add to the geometric com- associations: plexity of the spatial and temporal distribution of stars in Model A: origin as expanding dense embedded clusters giant OB associations [e.g., NGC 206 in M31 and NGC 604 Model B: origin in unbound turbulent giant molecular in M33 (Maíz-Apellániz et al., 2004)]. clouds A problem common to both models is that neither one We will label Model A as the KLL model (Kroupa et explains very well the spatial and temporal structure of OB al., 2001; Lada and Lada, 2003), while Model B will be associations, i.e., the fact that the subgroups seem to form named the BBC model (for Bonnell, Bate, and Clark) (see an age sequence (well known in Orion OB1 and Sco OB2). also Clark et al., 2005). One way out would be to argue that the subgroups are not The jury is still out on which model fits the observations always causally connected and did not originate by sequen- better. It may be that both models contain elements of truth tial star formation as in the Elmegreen and Lada (1977) and thus are not mutually exclusive. Future astrometric paradigm. Star formation in clouds with supersonic turbu- surveys, Gaia in particular, will provide constraints to per- lence occurring in convergent flows may be of a more ran- form sensitive tests on these models, but that will not hap- dom nature and only mimic a causal sequence of trigger- pen for about another decade. We refer the reader to the ing events [e.g., in Sco OB2, the UCL and LCC subgroups, review of Brown et al. (1999), who also points out the prob- with ages of ~15 and ~17 m.y. from the D’Antona and Maz- lems involved in the definition of an OB association and the zitelli (1994) tracks, have hardly an age difference at all]. division of these vast stellar aggregates into subgroups (see This is a problem for the KLL model, as sequential star Brown et al., 1997). formation can hardly generate two adjacent clusters (sub- We now discuss both models in turn. The KLL model groups) within such a short timespan. It is also a problem proposes that ultimately the star-formation efficiency in for the BBC model, but for a different reason. The fact that most embedded, incipient star clusters is too low (<5%) for star formation must be rapid in unbound transient molecu- these clusters to remain bound (see references in Elmegreen lar clouds is in conflict with the ages of the Sco OB2 sub- et al., 2000; Lada and Lada, 2003) after the massive stars groups (1, 5, 15, and 17 m.y.), if all subgroups formed from have expelled the bulk of the lower-density, leftover clus- a single coherent (long-lived) GMC. Still, the observational ter gas that was not turned into stars, or did not get accreted. evidence suggesting that Upper Sco can be understood in Therefore the cluster will find itself globally unbound, al- the context of triggered star formation (see section 7.2) can though the cluster core may remain bound (Kroupa et al., be reconciled with these models if we consider a scenario 2001) and survive. Such an originally compact but expand- with multiple star-formation sites in a turbulent large GMC, ing cluster could evolve into an extended subgroup of an in which triggering may easily take place. It remains to be OB association after a few million years. Essentially, the seen if subgroups of OB associations had some elongated KLL model describes how the dense gas in a bound and minimum size configuration, as Blaauw (1991) surmised, virialized GMC is partly converted into stars, most of which on the order of 20 pc × 40 pc. If so, OB associations are are born in dense embedded clusters, the majority of which then something fundamentally different from embedded disperse quickly (so-called cluster “infant mortality”). clusters, which would have many ramifications for the ori- The BBC model supposes that GMCs need not be re- gin of OB stars [However, a caveat with tracing back mini- garded as objects in virial equilibrium, or even bound, for mum size configurations (Brown et al., 1997) is that, even them to be sites of star formation. Globally unbound or mar- using modern Hipparcos proper motions, there is a tendency ginally bound GMC can form stellar groups or clusters very to obtain overestimated dimensions because present proper quickly, over roughly their crossing time (cf. Elmegreen et motions cannot resolve the small velocity dispersion. This al., 2000) [recent simulations also show that GMC them- situation should improve with Gaia, and with the newer cen- 356 Protostars and Planets V sus of low-mass stars, that can potentially provide statisti- 2001; Briceño et al., 2001, 2005a; Briceño et al., in prepa- cally robust samples to trace the past kinematics of these ration), and Cepheus (Sicilia-Aguilar et al., 2005), show that regions.] the groupings of stars with ages >4–5 m.y. have mostly lost their natal gas. The growing notion is that not only do mo- 7. CONSTRAINTS ON RAPID AND lecular clouds form stars rapidly, but that they are transient TRIGGERED/SEQUENTIAL STAR FORMATION structures, dissipating quickly after the onset of star formation. This dispersal seems to be effective in both low-mass 7.1. The Duration of Star Formation regions as well as in GMC complexes that give birth to OB associations. The problem of accumulating and then dis- One of the problems directly related to the properties of sipating the gas quickly in molecular clouds has been ad- the stellar populations in OB associations are the lifetimes dressed by Hartmann et al. (2001). The energy input from of molecular clouds. Age estimates for GMCs can be very stellar winds of massive stars, or more easily from SN discordant, ranging from ~108 yr (Scoville et al., 1979; Sco- shocks, seems to be able to account for the dispersal of the ville and Hersh, 1979) to just a few 106 yr (e.g., Elmegreen gas on short timescales in the high-density regions typical et al., 2000; Hartmann et al., 2001; Clark et al., 2005). of GMC complexes, as well as in low-density regions, like Molecular cloud lifetimes bear importantly on the picture Taurus or Lupus; if enough stellar energy is input into the we have of the process of star formation. Two main views gas such that the column density is reduced by factors of have been contending among the scientific community dur- only 2–3, the shielding could be reduced enough to allow ing the past few years. In the standard picture of star for- dissociation of much of the gas into atomic phase, effec- mation, magnetic fields are a major support mechanism for tively “dissipating” the molecular cloud. clouds (Shu et al., 1987). Because of this, the cloud must somehow reduce its magnetic flux per unit mass if it is to 7.2. Sequential and Triggered Star Formation attain the critical value for collapse. One way to do this is through ambipolar diffusion, in which the gravitational force Preibisch et al. (2002) investigated the star-formation pulls mass through the resisting magnetic field, effectively history in Upper Sco. A very important aspect in this con- concentrating the cloud and slowly “leaving the magnetic text is the spatial extent of the association and the corre- field behind.” From these arguments it follows that time- sponding crossing time. The bulk (70%) of the Hipparcos scale for star formation should be on the order of the dif- members (and thus also the low-mass stars) lie within an fusion time of the magnetic field, t ~ 5 × 1013 (n /n ) yr area of 11° diameter on the sky, which implies a character- D i H2 (Hartmann, 1998), which will be important only if the ion- istic size of the association of 28 pc. They estimated that the ization inside the cloud is low (n /n < 10–7), in which case original size of the association was probably about 25 pc. i H2 7 tD ~ 10 yr; therefore, the so called “standard” picture de- de Bruijne (1999) showed that the internal velocity disper- picts star formation as a “slow” process. This leads to rele- sion of the Hipparcos members of Upper Sco is only 1.3 km vant observational consequences. If molecular clouds live s–1. This implies a lateral crossing time of 25 pc/1.3 km s–1 for long periods before the onset of star formation, then we ~20 m.y. It is obvious that the lateral crossing time is much should expect to find a majority of starless dark clouds; (about an order of magnitude) larger than the age spread however, the observational evidence points to quite the con- of the association members [which is <2 m.y. as derived by trary. Almost all cloud complexes within ~500 pc exhibit Preibisch and Zinnecker (1999)]. This finding clearly shows active star formation, harboring stellar populations with ages that some external agent is required to have coordinated the ~1–10 m.y. Another implication of “slow” star formation onset of the star-formation process over the full spatial ex- is that of age spreads in star-forming regions. If clouds such tent of the association. In order to account for the small as Taurus last for tens of millions of years there should exist spread of stellar ages, the triggering agent must have crossed a population of PMS sequence stars with comparable ages the initial cloud with a velocity of at least ~15–25 km s–1. (the “post-T Tauri problem”). Many searches for such “mis- Finally, some mechanism must have terminated the star- sing population” were conducted in the optical and in X- formation process at most about 1 m.y. after it started. Both rays, in Taurus, and in other regions (see Neuhäuser, 1997). effects can be attributed to the influence of massive stars. The early claims by these studies of the detection of large In their immediate surroundings, massive stars generally numbers of older T Tauri stars widely spread across sev- have a destructive effect on their environment; they can eral nearby star-forming regions were countered by Briceño disrupt molecular clouds very quickly and therefore prevent et al. (1997), who showed that these samples were com- further star formation. At somewhat larger distances, how- posed of an admixture of young, X-ray active ZAMS field ever, the wind- and supernova-driven shock waves originat- stars and some true PMS sequence members of these re- ing from massive stars can have a constructive rather than gions. Subsequent high-resolution spectroscopy confirmed destructive effect by driving molecular cloud cores into col- this idea. As discussed in section 4, presently there is little lapse. Several numerical studies (e.g., Boss, 1995; Foster evidence for the presence of substantial numbers of older and Boss, 1996; Vanhala and Cameron, 1998; Fukuda and PMS stars in and around molecular clouds. Hanawa, 2000) have found that the outcome of the impact Recent widefield optical studies in OB associations like of a shock wave on a cloud core mainly depends on the type Sco-Cen (Preibisch et al., 2002), Orion (Dolan and Mathieu, of the shock and its velocity: In its initial, adiabatic phase, Briceño et al.: Low-mass Populations in OB Associations 357

the shock wave is likely to destroy ambient clouds; the later, Other relatively nearby regions have also been suggested isothermal phase, however, is capable of triggering cloud as scenarios for triggered star formation. In Cepheus, a collapse if the velocity is in the right range. Shocks travel- large-scale ring-like feature with a diameter of 120 pc has ing faster than about 50 km s–1 shred cloud cores to pieces, been known since the time of the Hα photographic atlases while shocks with velocities slower than about 15 km s–1 of H II regions (Sivan, 1974). Kun et al. (1987) first iden- usually cause only a slight temporary compression of cloud tified the infrared emission of this structure in IRAS 60- cores. Shock waves with velocities in the range of ~15– and 100-µm sky flux maps. The Cepheus bubble includes 45 km s–1, however, are able to induce collapse of molecular the Cepheus OB2 association (Cep OB2), which is partly cloud cores. A good source of shock waves with velocities made up of the Tr 37 and NGC 7160 open clusters, and in that range are supernova explosions at a distance between includes the H II region IC 1396. Patel et al. (1995, 1998) ~10 pc and ~100 pc. Other potential sources of such shock mapped ~100 deg2 in Cepheus in the J = 1–0 transition of waves include wind-blown bubbles and expanding H II re- CO and 13CO. Their observations reveal that the molecular gions. Observational evidence for star-forming events trig- clouds are undergoing an asymmetrical expansion away gered by shock waves from massive stars has been discussed from the galactic plane. They propose a scenario in which in, e.g., Carpenter et al. (2000), Walborn et al. (1999), the large-scale bubble was blown away by stellar winds and Yamaguchi et al. (2001), Efremov and Elmegreen (1998), photoionization from the first generation of OB stars, which Oey and Massey (1995), and Oey et al. (2005). are no longer present (having exploded as supernovae). The For the star burst in Upper Sco, a very suitable trigger ~10-m.y.-old (Sicilia-Aguilar et al., 2004) NGC 7160 clus- is a supernova explosion in the Upper Centaurus-Lupus ter and evolved stars such as µ Cephei, VV Cephei, and association that happened about 12 m.y. ago. The structure ν Cephei are the present-day companions of those first OB and kinematics of the large H I loops surrounding the Scor- stars. Patel et al. (1998) show that the expanding shell be- pius-Centaurus association suggest that this shock wave comes unstable at ~7 m.y. after the birth of the first OB passed through the former Upper Sco molecular cloud just stars. The estimated radius of the shell at that time (~30 pc) about 5–6 m.y. ago (de Geus, 1992). This point in time is consistent with the present radius of the ring of O- and agrees very well with the ages found for the low-mass stars B-type stars that constitute Cep OB2. Within a factor of <2, as well as the high-mass stars in Upper Sco, which have this age is also consistent with the estimated age for the been determined above in an absolutely independent way. Tr 37 cluster (~4 m.y.) (Sicilia-Aguilar et al., 2004). Once Furthermore, since the distance from Upper Centaurus- the second generation of massive stars formed, they started Lupus to Upper Sco is about 60 pc, this shock wave prob- affecting the dense gas in the remaining parent shell. The ably had precisely the properties (v ~ 20–25 km s–1) that gas around these O stars expanded in rings like the one seen are required to induce star formation according to the mod- in IC 1396. The dynamical timescale for this expansion is on eling results mentioned above. Thus, the assumption that the order of 1–3 m.y., consistent with the very young ages this supernova shock wave triggered the star-formation pro- (~1–2 m.y.) of the low-mass stars in the vicinity of IC 1396 cess in Upper Sco provides a self-consistent explanation of (Sicilia-Aguilar et al., 2004). This H II region is interpreted all observational data. as the most recent generation of stars in Cep OB2. The shock-wave crossing Upper Sco initiated the forma- In the Orion OB1 association Blaauw (1964) proposed tion of some 2500 stars, including 10 massive stars upward that the ONC is the most recent event in a series of star- of 10 M . When the newborn massive stars “turned on,” forming episodes within this association. The increasing they immediately started to destroy the cloud from inside ages between the ONC, Ori OB1b, and Ori OB1a have been by their ionizing radiation and their strong winds. This af- suggested to be a case for sequential star formation (Blaauw, fected the cloud so strongly that after a period of <1 m.y. the 1991). However, until now it has been difficult to investi- star-formation process was terminated, probably simply be- gate triggered star formation in Ori OB1 because of the lack cause all the remaining dense cloud material was disrupted. of an unbiased census of the low-mass stars over the entire This explains the narrow age distribution and why only region. This situation is changing with the new large-scale about 2% of the original cloud mass was transformed into surveys (e.g., Briceño et al., 2005a) that are mapping the stars. About 1.5 m.y. ago the most massive star in Upper low-mass population of Ori OB1 over tens of square de- Sco, probably the progenitor of the pulsar PSR J1932+1059, grees; we may soon be able to test if Orion can also be inter- exploded as a supernova. This explosion created a strong preted as a case of induced, sequential star formation. shock wave, which fully dispersed the Upper Sco molecular cloud and removed basically all the remaining diffuse 8. CONCLUDING REMARKS material. It is interesting to note that this shock wave must have Low-mass stars (0.1 < M < 1 M ) in OB associations are crossed the ρ Oph cloud within the last 1 m.y. (de Geus, essential for understanding many of the most fundamental 1992). The strong star-formation activity we witness right problems in star formation, and important progress has been now in the ρ Oph cloud might therefore be triggered by this made during the past years by mapping and characterizing shock wave (see Motte et al., 1998) and would represent the these objects. 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